A&A 503, 929–944 (2009) Astronomy DOI: 10.1051/0004-6361/200811349 & c ESO 2009 Astrophysics

Observation and modelling of star VIII. High resolution observations of M and K dwarf chromospheric lines

E. R. Houdebine1, K. Junghans2,M.C.Heanue3, and A. D. Andrews2

1 25 Rue du Dr. Laulaigne, 49670 Valanjou, France e-mail: [email protected] 2 Armagh Observatory, College Hill, BT61 9DG Armagh, N. Ireland 3 Dublin Institute for Advanced Studies, Dunsink Observatory, Castleknock, Dublin 15, Ireland

Received 14 November 2008 / Accepted 10 May 2009

ABSTRACT

Aims. We report on high resolution observations of chromospheric lines for M and K dwarfs. Methods. We observed up to 13 spectral lines per star: the Hα,Hβ,H ,Pa ,Pa8, CaII H, CaII IR triplet, CaI 6572Å, NaI D1/D2 and HeI D3 lines. We observed two dMe stars, one dM(e) star, 5 dM stars and 3 dK stars. Results. We observed a self-reversal in the emission core of the CaII H line for the brightest stars only, indicating a rather optically thick region of formation. We present original spectra of the NaI doublet and the CaII IR triplet for active dMe stars and less active dM stars. Core emission is detected in the Sodium lines and the CaII IR triplet lines for the most active M dwarf AU Mic. We investigate the difference spectra between active dMe stars and dM stars and show that these provide interesting new constraints for the NLTE-radiation transfer modelling of the chromospheres. In our sample, emission Hα profiles have a rather homogeneous FWHM of about 1.5 Å. This, according to our previous modelling, can be interpreted as the signature of a rather constant temperature break in the . We found that one of our targets (MCC 332) is a binary with a faint but active Hα emission component. ForthefirsttimewedetectthePa line for six dwarfs. It appears as weak absorption with possible weak wing emission in AU Mic. The region of the Pa8 line was observed but the line was not detected. Key words. stars: late-type – stars: activity – stars: chromospheres – stars: magnetic fields

1. Introduction (e.g. Houdebine & Stempels 1997, hereafter Paper VI). The FWHM of these lines is known to correlate with the stellar lu- High levels of magnetic activity in main sequence stars have minosity, known as the Wilson-Bappu relationship (e.g. Wilson manifested themselves as flares (e.g. Cristaldi & Rodonó 1970) & Bappu 1957, Elgaroy et al. 1999), and appears to be a diag- and chromospheric emission lines (e.g. Vaughan et al. 1981; nostic of the chromospheric mass column density (Ayres 1979). Giampapa et al. 1982; Rutten et al. 1989). Such phenomena are The CaII K to CaII H flux ratio seems a good diagnostic of the also observed in pre-main sequence stars (e.g. Shara et al. 1997) mid-chromosphere optical depth in M dwarfs (Paper VI). Hα is and RS CVn systems (e.g. Dempsey et al. 2001; Osten et al. often used to detect the highest activity stars (e.g. Worden et al. 2003). The most notable chromospheric lines, that serve as stel- 1981; Herbst & Layden 1987;Hawleyetal.1996)andasadi- lar classification criteria are the Balmer Hα line and the CaII agnostic of the upper chromosphere (e.g. Cram & Mullan 1979, resonance doublet. Stars showing Hα in emission were classi- 1985; Houdebine & Doyle 1994a,b, hereafter Papers I and II; fied as emission line stars. For the most active red dwarfs, the Houdebine et al. 1995, hereafter Paper III). chromospheric density is so high that it drives a number of lines There has been several efforts to model the spectral signa- into emission (e.g. Cram & Mullan 1979;Giampapaetal.1982; ff tures of dM and dK stars. Some pioneering work was carried out Cram&Mullan1985)andalsoa ects the formation of usu- byCram&Mullan(1979): they found important properties of ally photospheric lines (e.g. Houdebine & Doyle 1995, hereafter the chromosphere of dM stars, namely, the Hα line is weak when Paper IV; Houdebine 2009a, hereafter Paper X). the chromospheric pressure is low, then the absorption increases High resolution observations of chromospheric lines from as the pressure increases, and the line eventually goes into emis- the near-UV to IR wavelengths are essential to understand sion for the largest pressures. They also found that in dMe and constrain the chromospheric phenomena. Such observations stars, the source functions of the Balmer lines are collisionally were used to constrain the properties of the chromosphere by controlled, whereas in lower pressure dM stars they are photo- means of NLTE-radiative transfer modelling in main-sequence ionization controlled. Kelch et al. (1979) attempted to model Hα stars. For instance, the EW of the CaII resonance lines are observations of EQ Vir and 61 Cyg B with semi-empirical mod- often used as a diagnostic of the level of magnetic activity els. The first real proper fits to observations were carried out by Giampapa, Worden & Linsky (1982). They managed to re-  Based on observations collected at the European Southern produce rather well the CaII K line observations of three dMe Observatory, La Silla, Chile. and two dM stars. They found that the lower chromospheric

Article published by EDP Sciences 930 E. R. Houdebine et al.: Observations of M and K dwarf chromospheric lines temperature gradient is similar for all the M dwarf stars. Cram Walkowicz & Hawley (2008) analysed the chromospheric & Mullan (1985) found that strong Hα absorption profiles were spectral signatures of many dM3 stars, in a similar way indeed evidence for the presence of chromospheres in M dwarfs. as in Houdebine & Stempels (1997). Walkowicz (2008)and Cram & Giampapa (1987) calculated the Hα and CaII line pro- Walkowicz & Hawley (2009) derived models for dM3 stars. files assuming a slab chromosphere with varying parameters. They found that for low and intermediate activity stars, only a They found the interesting relation between the Hα equivalent single component model chromosphere is necessary to repro- width and the CaII K equivalent width, that was later confirmed duce observations, but that active dM3 stars necessitate a two- by calculations of grids of model chromospheres. component model chromosphere. Houdebine (2009a, Paper X) Subsequently, a substantial amount of work was carried proposed the first two-component model chromosphere for a out by the Armagh group (Houdebine 1990; Panagi 1990; dM1 star: namely the intermediate activity star Gl 205. Houdebine & Panagi 1990; Houdebine & Doyle 1994a,b; Doyle Some interesting studies of dK stars have also been per- et al. 1994; Houdebine & Doyle 1995a; Houdebine et al. formed: Cuntz et al. (1999) propose theoretical ab-initio mod- 1995; Houdebine & Doyle 1995b; Lanzafame 1995; Lanzafame els with two components. This was the first study to really & Byrne 1995; Houdebine et al. 1996; Short et al. 1997; propose two-component solutions for the observations: the qui- Andretta et al. 1997; Short & Doyle 1998a,b; Jevremovic et al. escent region component is heated by acoustic waves, and the 2000). Houdebine & Doyle (1994a,b, Papers I and II there- active region component, magnetic in character, is heated by lon- after) studied in detail the modelling of the Balmer lines and gitudinal tube waves. Sim & Jordan (2005) developed a model the Lyman/Balmer flux ratio in the dM1e star AU Mic. They for  Eri (dK2) based on the emission measure distribution of proposed for the first time a model that could reproduce these UV lines. They extend their model to lower chromospheric tem- observations. peratures in a semi-empirical way. They also derive a crude two- Doyle et al. (1994) and Houdebine & Doyle (1995b) pro- component model for this star. posed chromospheric models for the other end of the activity Recently, in Paper XI in this series (Houdebine 2009b) we range: the basal chromosphere M dwarfs for which the contri- defined a new method to constrain the theoretical profiles of the bution from plages can be neglected. On the basis of the models Hα line and the CaII H & K resonance lines from the measure- of the two extremes of the activity range, dM1e stars and dM1 ments of their equivalent widths. This method is based on the basal stars, Houdebine et al. (1995) built the first grid of model calculations of the grid of models of Paper VI. We showed that, atmospheres that cover the whole range of activity for dM1 stars. providing that such a grid of model chromospheres is available, Lanzafame (1995) and Lanzafame & Byrne (1995) analyzed one can derive two-component model chromospheres for quies- several spectral lines of the dMe star Gl 182. Lanzafame (1995) cent and active regions on dM1 stars. In Paper XI, we derived produced two grids of model chromospheres with a high col- such model chromospheres for nine near-solar dM1 umn mass at the transition region between log(M) = −4.0to stars. Many were found to be rather active, with sometimes ac- log(M) = −4.15. They could reproduce partly the Hα line profile tivity levels close to those found in dM1e stars. Such models are but the self-reversal of the models was too strong. very important because they allow us to understand the proper- Shortetal.(1997) focused their analysis on the CaI 4227 Å ties of quiescent regions as well as active regions on M dwarfs. line. Andretta et al. (1997) investigated the formation of the NaI Houdebine (2009c, hereafter Paper XII) also developed two- D1/D2 resonance lines in dM stars. They found that these lines component model chromospheres with the same technique for are interesting chromospheric diagnostics and that it is neces- seven dM1e stars. sary to include a proper treatment of the background opacities. There are several other spectral lines of great interest such as / Short & Doyle (1998a) modelled the Hα line and the Paβ line in the NaI D1 D2 lines, the CaII IR triplet lines and a large number the dM1e star AU Mic. They could reproduce separately, but not of faint emission lines in the near-UV (Fuhrmeister et al. 2005). simultaneously the Hα line and the Paβ line. Again their mod- These lines can constrain the chromosphere from the tempera- els assumed a uniform chromosphere. Short & Doyle (1998b) ture minimum to the transition region. However, most observa- attempted to model simultaneously the Hα line and the NaI D tions were dedicated to the study of the CaII resonance lines or lines. The profiles cannot be simultaneously fitted for all but the Hα line. Here, we attempt to gather observations of as many one of their stars. Jevremovic et al. (2000) studied the effects lines as possible for a few stars. of micro-turbulent velocities on the formation of the Hα line, the Observations of spectral lines are also the basis for devel- CaII resonance lines and the NaI resonance lines. They found oping NLTE-radiation transfer model chromospheres. The chro- that this chromospheric parameter has an important influence on mospheres of M dwarfs is still little understood in particular the formation of these lines. for intermediate activity stars, although some recent work has Some other authors have also contributed to the progress of developed two-component model chromospheres for these stars semi-empirical modelling of dM stellar chromospheres. Mauas (Papers XI and XII). In Paper VI, we built a grid of model chro- & Falchi (1994) modelled the chromosphere of AD Leonis as- mospheres for M1 dwarfs (including AU Mic) and calculated suming a homogeneous chromosphere. Mauas et al. (1997)de- the Hα and CaII H & K line profiles, that we successfully com- veloped models for two dM stars considered as basal, based pared to high resolution observations. A number of authors built on the fit of several lines. Houdebine & Stempels (1997)pro- some grids of model atmospheres, but exact models of single posed an improved grid of 30 model chromospheres for dM1 stars are few (see references above). This kind of model is es- stars, from basal to the most active dM1e stars. Baranovskii et al. sential to give constraints to the grids of model atmospheres (e.g. (2001) modelled three dMe stars on the basis of Balmer line ob- Paper VI). Both detailed modelling and grids of models are nec- servations. Busa et al. (2001) re-investigated in detail the im- essary to understand the chromospheric phenomenon. Detailed portance of line blanketing in the formation of several impor- models give constraints to the grids of models, and from the grids tant chromospheric lines in the case of main-sequence dwarfs. of models we can build two-component models (Papers XI and Andretta et al. (2005) studied the CaII infrared triplet as a po- XII). tential stellar activity diagnostic. Osten et al. (2006) derived a Here, we gathered observations of as many lines as possible model atmosphere above 10 000 K for the dMe star EV Lac. for a few stars, of spectral types late dK and dM, in order to have E. R. Houdebine et al.: Observations of M and K dwarf chromospheric lines 931 the most complete constraints for the purpose of NLTE radiation 3. Results transfer modelling of the chromosphere. A partial report of these ff observations can be found in Houdebine (1990). However, this In total, we obtained 71 high resolution spectra for 10 di erent later study did not analyze nor comment on the results. The re- stars of spectral types late dK and dM. The summary of our ob- sults for the Paschen lines were not presented either. servations is given in Table 1. The spectra are shown from Fig. 1 to 9: CaII H & H in Fig. 1,Hβ in Fig. 2,NaID1&D2inFig.3, Hα in Fig. 4, CaII 8498 Å in Fig. 5, CaII 8543 Å in Fig. 6,CaII 2. Observations and data reduction 8662 Å in Fig. 7,Pa in Fig. 8 and Pa8 in Fig. 9.Wealsoshow the difference spectra AU Mic-Gl 887 in Fig. 10 and AU Mic- The observations presented below were collected during an ob- Gl 908 in Fig. 11. In our diagrams, all spectra were normalized serving run at the European Southern Observatory with the 1.4 m to the continuum (we define the continuum here as the upper Coudé Auxiliary Telescope equipped with the Coudé Echelle boundary delineated by the photospheric lines). A constant shift Spectrograph (CES) and a CCD detector. The resolving power was added in the Y-axis for better clarity. We show the zero level was 40 000. Simultaneous photometry was carried out with the for each spectrum as a dotted-dashed line. We discuss below our 50 cm telescope and a standard photometer in order to assess observations and the properties of the chromosphere for each the flaring or quiescent state of the stars during observations. star separately. Weather conditions at ESO varied from poor to good and as a We present here the first observations and detections of the consequence, telluric water vapor lines can be strong. Pa line in 6 stars (Fig. 8). The region of Pa8 was also observed Our aim was to collect the most complete set of chromo- for two stars but this line was not detected (Fig. 9).Wealso spheric line profiles in high resolution for a few stars with dif- present original observations of the NaI doublet (Fig. 3)andthe ferent spectral types and activity levels. We also observed stars CaII IR triplet lines (Figs. 5 to 7). with similar spectral types but different levels of activity in order For the Pa line, the spectra are heavily contaminated by tel- to assess its effects on line profiles, from the photosphere to the  luric absorption features. We observed a standard star, HD 8946 transition region. These spectra are used as the basis for the de- (spectral type A0), in order to remove these telluric spectral fea- tailed NLTE modelling of their chromosphere. The Balmer line tures. The spectra were all normalized to the continuum, and the observations for AU Mic were used to develop a model chromo- spectrum of HD 8946 was subtracted from the spectra of our tar- sphere for this star (Papers I and II). Some results of the CaII get stars. This yields the line profile of Pa which is not present lines and Balmer lines for M1 dwarfs were also compared to the  in HD 8946. We then added a constant (unity) to the subtracted grid of models of Paper VI. We will use the data presented here spectra in order to mimic the background continuum and mea- to build a model chromosphere for each individual star. The stars sure the equivalent widths. The resulting spectra are shown in presented here span from low activity to high activity. The result- Fig. 8. ing detailed model chromospheres will be used in future studies ffi to develop new grids of model chromospheres. We emphasize It is di cult to measure spectral line equivalent widths for that this data is rather complete and can be used to constrain the late dK and dM stars because of the uncertainty in the level of atmosphere from the photosphere to the lower transition region. the adjacent continuum. Indeed, the background continuum con- Observations at ESO were obtained over five nights in differ- tains numerous faint absorption photospheric features. In order ent spectral regions. This is of the order of the rotation period of to measure the equivalent widths, we defined two regions 3 Å the most active dwarfs, which therefore may result in unwanted wide on each side of the spectral line. We averaged the flux in rotational modulation in the line profiles. Nevertheless, obser- these spectral regions that we define as the continuum level. The vations indicate so far that such modulation is somewhat weak spectral line flux is then integrated between these two regions. (e.g. Baliunas 1981). No flares were detected, either in white The equivalent width is the ratio of this flux to the average flux light or in the line profiles. We therefore assume that the spectra in the two spectral regions on either side of the line. discussed here are representative of the stellar quiescent state. The spectra were reduced following the standard procedure. 3.1. Gl 488 Dark exposures, flat-field and Thorium spectra were gathered regularly in order to assess variations in dark counts, flat-field Cenarro et al. (2006) classified this star as an M0.5 dwarf with irregularities and wavelength stability. We found that they were Teff = 3750 K, log(g) = 4.75 and [M/H] = 0.10. For Gl 488, quite stable in time. The spectra were flat-field corrected, sky we observed the Hα line, the Ca i 6572 and the Na i doublet. subtracted, and wavelength calibrated using the Thorium-Argon We measured an Hα line equivalent width of 0.588 ± 0.05 Å. spectra. The wavelength calibration gives typically an accuracy Wilson & Wooley (1970) recorded spectrograms of Gl 488 in of Δλ/λ ∼ 6 × 106 (i.e. ∼0.01 Å at 6500 Å). the Ca ii H and K lines. However they only performed visual However, we faced some problems when applying the flat- estimates of the intensity of these lines against a reference star field correction for wavelengths greater than 8000 Å, owing to 61 Cygni. The line intensity was +1 on their scale which means large amplitude interference fringes. Because the light paths of that it is slightly stronger than the reference 61 Cygni. Stauffer & the stellar and flat-field sources are not exactly identical in the Hartmann (1986) observed the Hα line and measured an equiv- spectrograph, the resulting fringe patterns also change slightly. alent width of 0.59 Å. This is the same as our measurement of This results in a slightly inadequate flat-field correction with an 0.588 ± 0.05 Å for this line. amplitude of about 5% of the normalised continuum. We also Mathioudakis & Doyle (1992) measured the flux in the emphasize that the fringe patterns are much broader than typical Mg ii h & k lines observed with IUE. They inferred a surface −1 −2 photospheric absorption lines. This implies that it does not af- flux of log FMgII = 5.64 ergs s cm . This is relatively high fect the detection of absorption lines, notably the Paschen lines. compared to other values inferred for dM stars and is compa- We generally found that a reasonably good correction can be rable to values observed for dMe stars. Similarly, Rutten et al. −1 −2 achieved (less than σ = 5% in the line profiles and continuum) (1989) report a surface flux log FCaII = 5.58 ergs s cm in that is sufficient for the purpose of this study. the Ca ii H & K lines. Wright et al. (2004) observed a CaII S 932 E. R. Houdebine et al.: Observations of M and K dwarf chromospheric lines

Table 1. Summary of ESO observations and line profile parameters.

Star Spectral UT Duration λ FWHM EW Photometry Parameters line start (s) Å Å mÅ U-band

Gl 488 Hα 23:57:35/21 1200 6563.532 0.95 588 ± 50 Y dM0.5 CaI 6572 23:06:35/21 1200 6573.56 – 281 ± 30 Y

v = 8.51 NaI D1,2 01:12:30/21 900 – – 11 780 ± 500 Y

Gl 526 CaII IRT 8542 23:54:04/19 350 – – 1920 ± 200 N dM4 v = 8.5

Gl 588 Hα 02:18:19/21 1800 6563.524 0.71 455 ± 50 Y

dM4 Hβ 23:03:21/23 4800 – – – N

v = 10.1H 02:30:59/23 2700 – – – Y CaI 6572 23:16:29/20 1800 6573.53 – 419 ± 30 Y CaII H 02:30:59/23 2700 ∗3968.94 ∗0.22 ∗−346 ± 10 N CaII IRT 8498 02:47:30/22 1800 8499.000 0.44 517 ± 50 Y CaII IRT 8542 01:30:23/20 900 8543.050 0.69 1234 ± 200 N CaII IRT 8662 04:59:16/22 1800 8663.179 0.72 1155 ± 200 Y

NaI D1,2 23:16:29/20 1800 – – 7330 ± 500 Y

Pa – – 9547.31 – 105 ± 10 –

Gl 664 Hα 02:59:33/21 600 6562.911 1.04 640 ± 50 Y

dK5 Hβ 05:50:21/23 900 – – ∼410 ± 50 N v = 6.34 CaI 6572 02:33:31/19 1340 6572.90 – 145 ± 10 N CaII IRT 8498 02:20:52/22 900 ∼8498.19 1.42 1070 ± 100 Y CaII IRT 8542 04:33:54/19 1260 8542.160 3.41 3030 ± 200 N CaII IRT 8662 05:34:24/22 900 8662.323 2.20 2120 ± 200 Y

NaI D1,2 02:33:31/19 1340 – – 6800 ± 500 N

Pa – – 9546.23 – 84 ± 10 –

Gl 729 Hα 03:14:42/21 1800 6562.392 1.23 −1230 ± 100 Y

dM4.5e Hβ 04:05:47/23 3000 4861.050 0.705 −1736 ± 100 N v = 10.6 CaI 6572 08:04:34/22 1800 6572.42 – 489, 5 ± 50 Y CaII IRT 8498 03:24:44/22 1800 ∼8497.58 – ∼385 ± 100 Y CaII IRT 8542 07:01:32/21 2400 8541.810 1.98 1340 ± 100 Y CaII IRT 8662 04:22:56/22 1800 8661.431 1.05 1055 ± 100 Y

NaI D1,2 08:04:34/22 1800 – – 8120 ± 500 Y

HeI D3 08:04:34/22 1800 ∼5875.35 – ∼−114 ± 10 Y

AU Mic Hα 04:57:26/21 900 6562.306 1.415 −2466 ± 100 Y

(Gl 803) Hβ 05:24:01/23 1200 4860.860 0.793 −1912 ± 100 N

dM1e H 07:59:32/23 1800 3969.98 ∼0.44 −1081 ± 100 N v = 8.61 CaI 6572 08:45:02/19 1200 6572.30 – 427 ± 50 N CaII H 07:59:32/23 1800 ∗3968.20 0.31 −4554 ± 200 N CaII IRT 8498 09:53:40/22 600 ∗8497.400 ∗0.45 ∗−381 ± 50 N 420±50 CaII IRT 8542 05:48:00/19 2700 ∗8541.429 ∗0.520 ∗−446 ± 50 N 1320 ± 200 ∗ NaI D1,2 09:16:42/19 1200 – – -62.9±10, −32.4 ± 10 N 7982 ± 500

HeI D3 09:16:42/19 1200 5875.271 0.551 129 ± 10 N

Pa – – 9545.34 – 173 ± 10 – E. R. Houdebine et al.: Observations of M and K dwarf chromospheric lines 933

Table 1. continued.

Star Spectral UT Duration λ FWHM EW Photometry Parameters line start (s) Å Å mÅ U-band

AX Mic Hα 04:01:37/21 600 6562.868 0.856 592 ± 50 Y

(Gl 825) Hβ 05:05:32/23 900 – – ∼200 ± 50 N

dK7 H 08:36:48/23 1200 – – ∼4.3 ± 2N v = 6.67 CaI 6572 09:45:21/19 900 6572.86 – 249 ± 30 N CaII H 08:36:48/23 1200 3968.515 0.33 −599 ± 50 N CaII IRT 8498 10:06:46/22 300 1.04 975 ± 100 N CaII IRT 8542 06:30:13/19 600 8542.83 2.13 2460 ± 200 N

NaI D1,2 09:45:21/19 900 – – 9111 ± 500 N

Pa – – 9546.12 – 92 ± 10 –

Gl 884 Hα 09:54:47/21 900 6562.576 0.98 605 ± 50 Y

dK5 H 10:28:04/23 1800 3969.93 – −1250 ± 100 N v = 7.89 CaI 6572 6572.55 – 235 ± 30 CaII H 10:28:04/23 1800 3968.33 0.32 −968 ± 50 N CaII IRT 8498 10:48:13/22 600 8497.907 1.56 1007 ± 100 Y

Pa – – 9545.63 – 46 ± 10 –

Gl 887 H 09:03:31/23 1200 – – ∼−110 ± 30 N dM0.5 CaII H 09:03:31/23 1200 3968.271 0.25 −796 ± 50 N v = 7.36 CaII IRT 8498 11:01:40/22 1030 8497.61 0.90 742 ± 100 N CaII IRT 8542 11:15:25/22 430 8541.63 2.07 2130 ± 100 N

Pa – – 9545.58 – 97 ± 20 –

Gl 908 Hα 09:32:34/21 900 6560.550 0.691 360 ± 50 Y dM1 CaI 6572 10:40:14/19 1300 6570.56 – 332 ± 50 N v = 8.98 CaII H 09:28:47/23 3300 3967.15 0.14 −188 ± 30 N CaII IRT 8498 10:28:55/22 900 8495.148 0.50 755 ± 100 N CaII IRT 8542 07:57:16/21 1200 8539.24 0.98 1950 ± 100 Y

NaI D1,2 10:40:14/19 1300 – – 7590 ± 500 N

MCC 332 Hα 05:19:10/21 1200 – – 288 ± 50 Y

dK7/dM3.5(e) Hβ 06:09:33/23 1800 – – ∼271 ± 40 N

v = 9.4H 06:52:24/23 3600 3969.98 – ∼−248 ± 30 N CaI 6572 07:13:19/22 1800 6572.29 – 197 ± 30 Y CaII H 06:52:24/23 3600 3968.22 0.31 −866 ± 50 N CaII IRT 8498 10:15:26/22 600 8497.30 1.93 ∼931 ± 200 Y CaII IRT 8542 06:25:03/21 1500 ∗8541.720 – ∗−57 ± 20 Y 2494 ± 200

NaI D1,2 07:13:19/22 1800 – – 6550 ± 500 Y

∗ Emission core only. The total exposure time and start of the observing sequence are given respectively in Cols. 4 and 3. The date of June 1989 is given after the Universal Time. In the last column we indicate when simultaneous photometry in the Johnson U-band was obtained with the ESO 50 cm telescope.

index of 1.793. The S -value is defined by the operation of the just below that of the dMe stars, when absorption in Hα attains Mount Wilson spectrometers, (Duncan et al. 1991) which mea- a maximum before going into emission (Houdebine & Stempels sure a quotient of the flux in two triangular bandpasses centered 1997). This agrees with the finding that this star is slightly metal on the H and K emission cores and two continuum regions on ei- rich (Paper VI). Indeed, in Paper VI we found that activity level ther side. The Hα equivalent width is also large for a dM0.5 star. correlates with radius in dM1 stars, and because radius correlates According to Houdebine & Stempels (1997; see also Houdebine with metallicity (Houdebine 2008, hereafter Paper VII) then ac- et al. 1995, Paper III) calculations, this shows that Gl 488 is an tivity correlates with metallicity. Therefore, metal rich stars are intermediate activity star. That is to say, it has an activity level on average more active than metal poor stars. 934 E. R. Houdebine et al.: Observations of M and K dwarf chromospheric lines

Fig. 2. We show the Hβ line for six stars.

observed the Ca ii H & K lines. They calculated excess sur- face fluxes (the flux from a radiative equilibrium atmosphere −1 −2 was subtracted) of log FH = 4.23 ergs s cm and log FK = 4.26 ergs s−1 cm−2 which is slightly smaller than the Blanco et al. (1974) flux. They also give the line widths: FWHMH = 0.26 Å and FWHMK = 0.27 Å. Cincunegui et al. (2007) found −1 −2 log FHK = 4.74 ergs s cm . Wright et al. (2004) report a CaII S index of only 0.787. Rauscher & Marcy (2006)giveCaIIH& K equivalent widths of 0.44 Å and 0.61 Å respectively. Here, we only observed the CaII 8542 Å line. This line has an average EW of 1.92 Å. Therefore, Gl 526 is a low activity slightly metal poor M1 dwarf. Gl 526 was also observed with IUE. Mathioudakis & Doyle (1989) gave a surface flux for the Mg ii h&klines;logFMg = 5.52, more than an order of magnitude larger than the flux in the Ca ii H & K lines. This departs from the nearly one to one cor- relation found by Panagi & Mathioudakis (1993) for these lines, and possibly arises from the variability of this star. Tripicchio et al. (1999) observed the Na i doublet and the K i λ7699 reso- Fig. 1. We show the CaII H line and the H line in increasing order of nance line and they gave equivalent widths of 0.85 Å and 8.7 Å the CaII H line equivalent width of the emission core. Note that the respectively. inner wings progressively disappear and that H is just about detectable Stauffer & Hartmann (1986) report an Hα equivalent width for McC 332, Gl 884 and Gl 887. of about 0.45 Å. Herbst & Layden (1987) give a very close value of 0.44 Å, determined photometrically. Kamper et al. (1997) 3.2. Gl 526 also measured an Hα equivalent width that is of the same order; 0.49 Å. Kraft et al. (1964) measured a FWHM of 0.74 Å for Hα. Gl 526 is an M1 dwarf with (R − I)c = 1.124, Teff = 3470 K, R = Cutispoto & Giampapa (1988) also observed spectroscopically / = − 0.582 R and [M H] 0.086 (Paper VII). It is therefore a small the Hα line for this star and although it is a relatively inactive M1 dwarf relative to solar metallicity M1 dwarfs (Paper VII) and star, they observed a day to day variation which they say is ev- is therefore slightly metal-poor. For Gl 526, we observed only idence for magnetic heating. Marcy & Chen (1992) measured a the Ca ii 8542 line. vsini for Gl 526 of 2.6±0.4kms−1, which is large for an inactive Wilson & Wooley (1970) observed the Ca ii H&Klines dM1 star compared to the average measurements of Paper VII, and report an activity index of –3, which implies a very low but remains significantly smaller than the values for dM1e stars activity compared to 61 Cygni. Blanco et al. (1974) observed (Marcy & Chen 1992; and Paper VII). the Ca ii K line and measured an equivalent width of the core When comparing the observations and models of Houdebine emission of 0.40 Å. From this they calculated a surface flux & Stencel (1997), Gl 526 appears to be a relatively inactive star. −1 −2 of log FK = 4.51 ergs s cm . Giampapa et al. (1981)also Giampapa, Worden & Linsky (1982) found a small plage filling E. R. Houdebine et al.: Observations of M and K dwarf chromospheric lines 935

Fig. 3. We show the NaI doublet lines in decreasing order of the equiv- alent width. The HeI D3 line is also visible on the left for emission line Fig. 4. We show the Hα line in decreasing order of the equivalent width. stars. Note the interesting P-Cygni profile for McC 332. factor of only 1.3%. This star has also been studied in more detail by Houdebine (2009b). They produced a two-component model The Hα EW is rather large; 0.46 ± 0.05 Å. This suggests that chromosphere for this star and found that this star has a rather Gl 588 is an intermediate activity star, although its CaII line low pressure chromosphere. Contrary to Giampapa et al. (1982), fluxes are relatively low; −0.35 ± 0.01 Å. More detailed mod- they found a plage filling factor of 36%. In general, they found elling is required to make conclusions regarding the activity level that the filling factor does not vary much with the activity level, of this star. but that it is the chromospheric pressure that diminish with the The NaI doublet line EW is not large, compared to Gl 729, a activity level. dM4.5e star. This comes from the fact that the line width is much smaller in Gl 588. Gl 588 has no detectable Hβ. It has also the 3.3. Gl 588 weakest and narrowest CaII IR triplet lines (besides dMe stars) which also indicates a small metallicity for this star. These lines Gl 588 is a dM4 star of v magnitude 10.1. For Gl 588, we ob- are unlikely to be filled-in because of the relatively weak activity served the Hα,Hβ,H,Caii H, Ca i 6572 lines, the Ca ii IR level of Gl 588, and because even Gl 729, a dM4.5e star does triplet lines, the Na i doublet and the Pa line. not show any core emission in these lines. Therefore, the weak 936 E. R. Houdebine et al.: Observations of M and K dwarf chromospheric lines

Fig. 5. We show the CaII IR triplet line at 8498 Å in decreasing order Fig. 6. We show the CaII IR triplet line at 8542 Å in decreasing order of the equivalent width. Note the emission core in AU Mic. Gl 729 has of the equivalent width. Note the emission core in AU Mic. Gl 729 has no core emission although it is a dMe star. no core emission although it is a dMe star.

nature of the CaII IR triplet lines in the spectrum of Gl 588 is probably due to low metallicity. This is further emphasized by They also gave an equivalent width for Hα of 0.45 Å, which the fact that these lines are also narrower than for other stars, is comparable to the value we report here; 0.455 ± 0.05 Å. a signature of low metallicity. Such an effect of metallicity was Mathioudakis & Doyle (1991) also observed Hα but at a lower also found for the NaI doublet lines in dM1 stars (Houdebine resolution (1.4 Å). They measured an Hα equivalent width of 2009d, Paper XIII). This would explain why the CaII resonance only 0.32 Å. This is significantly smaller than other observa- lines are weak whereas Hα is not. Pa was also detected which tions, and we believe this is due to the low resolution used. is consistent with the fact that Hα is relatively strong for a dM4 But this difference could also be due to a different choice in the star. level of the continuum which is difficult to determine for these Robinson et al. (1990) gave measurements of chromospheric stars. Mathioudakis & Doyle (1991) also measured the flux re- −1 −2 −14 −2 −1 radiative losses, of log FH = 3.89 ergs s cm and log FK = ceived at in the Ca ii K line; 4.1 × 10 erg cm s , −1 −2 −1 −2 4.01 ergs s cm respectively for the Ca ii H and K lines. which yields a surface flux of log FCaIIK = 4.33 ergs s cm . E. R. Houdebine et al.: Observations of M and K dwarf chromospheric lines 937

Fig. 7. We show the CaII IR triplet line at 8662 Å in decreasing order of the equivalent width. Gl 729 has no core emission although it is a dMe star.

Mathioudakis & Doyle (1992)giveasurfacefluxintheMgii −1 −2 lines of log FMgII = 4.98 ergs s cm which is higher than the flux in the Ca ii lines.

3.4. Gl 664 Gl 664 is a dK5 star. It belongs to the nearby triple system 36 Ophiuchi (Gl 664 = 36 Ophiuchi C). An attempt to deter- mine the of A and B was performed by Brosche (1960). He found two possibilities characterized by a = 25.5 and P = 11547 yr, and a = 13.8 and P = 548.7 yr. The second solution gives a mass of 0.73 M for each star in the system for an eccentricity of the ≤0.90 (Cayrel de Strobel et al. 1989). The distance of the C component to the AB components is much larger and its orbit cannot be determined. Cenarro et al. (2006) found that for this star the spectral type is dK5, Teff = 4540 K, log(g) = 4.54 and [M/H] = −0.37. Diaz et al. (2007)alsogive [M/H] = −0.21. This gives an estimate of the error on this mea- surement. It is therefore metal poor. Fig. 8. We show here the Pa line at 9545 Å. In the upper panel,we For Gl 664, we observed the H ,H,Cai 6572 lines, the  i i α β show the normalized spectrum for AU Mic, together with the spectrum Ca IR triplet lines, the Na doublet and the Pa line. Cayrel of AU Mic minus that of HD 8946 (an A0 star). In the lower panel, de Strobel et al. (1989) also observed the Ca ii 8498.06 and we show the spectra of our targets to which has been subtracted the Ca ii 8542.14 IR triplet lines at high resolution. They mea- spectrum of HD 8946, plus the constant 1. sured a central depth of 0.46 and 0.61 (respective to the adjacent continuum) respectively for the 8498.06 and the 8542.14 lines. Early high resolution observations of Hα byKraftetal. (1964)gaveaFWHM for this line of 0.91 Å, which is close Baliunas et al. (1995) report on regular observations in the to the value we observe here of 1.04 Å. Robinson et al. (1990) ii ii Ca H and K lines (narrow band photometry) since 1966. They observed Gl 664 at high resolution in the Hα and Ca H measured an average index S of 0.770 and also found an activity and K lines. They calculated a “chromospheric flux” (flux in cycle period of 21 . Henry et al. (1996) observed more than the emission line core) in the H and K lines of respectively 800 southern stars in the region of the Ca ii H and K lines with = −2 −1 = −2 −1 log FH 5.16 erg cm s and log FK 5.28 erg cm s . a medium resolution spectrograph. They inferred an S index of They also measured an Hα equivalent width of 750 mÅ. Herbst 1.07, somewhat larger than Baliunas et al.’ (1995). Wright et al. & Miller (1989) performed some Hα photometry and measured (2004) report an S index of 0.858 a bit higher than Baliunas an Hα equivalent width of 850 mÅ. Here we report a signifi- et al.’ (1995). Cincunegui et al. (2007)giveameanSindexof cantly smaller equivalent width of 650 ± 50 mÅ. 0.76. The mean of all measures is S = 0.865 which is relatively 938 E. R. Houdebine et al.: Observations of M and K dwarf chromospheric lines

3.5. Gl 729

Gl 729 is a dM4.5 Hα emission line star. For Gl 729, we observed the Hα,Hβ,Cai 6572 lines, the Ca ii IR triplet lines, the Na i doublet and the He i D3 line. Young et al. (1989) report on Hα observations of Gl 729. They give an “excess equivalent width” of 1.72 Å for Gl 729. This “excess equivalent width” is defined as the measured equiv- alent width minus the equivalent width of the lower envelope of the graph given by Stauffer & Hartmann (1986). However, ac- cording to calculations by Houdebine et al. (Papers III and VI), the maximum absorption in Hα corresponds to intermediate ac- tivity levels which cannot be taken as a reference for minimum radiative losses. We recovered the original equivalent widths from the relation they used to calculate their “excess equivalent width”. We find Hα equivalent widths of –2.14 Å, –1.39 Å and –1.35 Å, which are of the same order as the value we measured; –1.23 Å. Herbst & Layden (1987) measured photometrically an Hα equivalent width of –1.11 Å, which is again close to our mea- surement. Panagi & Mathioudakis (1993)giveasurfacefluxinthe Ca ii H & K lines of log F = 5.95 erg cm−2 s−1,whichis Fig. 9. We show the region of the Pa8 line at 8662 Å, for AU Mic and high, and higher than for AU Mic which has a larger Hα equiv- AU Mic-AX Mic. There seems to be no real detection of this line at alent width. Rauscher & Marcy (2006) measured EW of the H 9015 Å in our spectra. and K lines of –3.93 Å and –5.82 Å respectively. Jenkins et al.  =  (2006) found log RHK 4.29 only (RHK is a measure of the chromospheric contribution to the H & K lines, see Noyes et al. 1984). This is in significant disagreement with the Panagi & Mathioudakis (1993) measurement. Landsman & Simon (1993) −12 −2 −1 high (all these measures are in the Mount Wilson S index). The measured a Lyα flux at Earth of 1.40 × 10 erg cm s , which taking into account the distance of the star yields a flux of variations, from 0.76 to 1.07 are due to the variability of the star − − (Baliunas et al. 1995). 6.6 × 105 erg cm 2 s 1 on the stellar surface. Doyle et al. (1990) give a slightly higher surface flux of 8.5 × 105 erg cm−2 s−1,and The He i D3 line (5876 Å) was observed at high resolution − − also a Mg ii h & k flux of 4.0 × 105 erg cm 2 s 1. for Gl 664 by Saar et al. (1997). They observed an equivalent i width for this line of 11 ± 4 mÅ. Although this value is weak, it Johns-Krull & Valenti (1996) observed the Fe 8468.40 Å is evidence for a chromosphere. Tripicchio et al. (1999) observed line. From their observations they inferred that 50% of Gl 729 is . ± . the K i λ7699 resonance line at high resolution and measured an covered with magnetic fields of 2 6 0 3 kG strength. equivalent width of 460 mÅ. Gl 664 was also observed in the Mg ii k line with IUE: Scoville & Mena-Werth (1998) measured 3.6. AU Mic (gl 803) a FWHM for the Mg ii k line of 0.39 Å. AU Mic is one of the best studied active red dwarfs, it is also a Baliunas et al. (1983) observed a rotation period for this star flare star (e.g. Kunkel 1973). It is known to have a disk of debris ± of 18.0 0.5 days. Saar and Osten (1997) investigated rota- (e.g. Fitzgerald et al. 2007). For AU Mic, we gathered one of tion and turbulence in a number of late type dwarfs and found the most complete set of observations: the Hα,Hβ,H,Caii H, for Gl 644 that v sin i and vmac (the macroturbulent velocity) are − − Ca i 6572 lines, the Ca ii IR 8498 Å and 8542 Å triplet lines, the 0.3 km s 1 and 1.5 km s 1 respectively. Fekel (1997) also cal- − Na i doublet, the He i D3, the Pa and Pa lines. culated these parameters and found values of 2.2 km s 1 and  8 2.0 km s−1 respectively. Saar & Osten (1997) also inferred a ro- A number of studies have been performed on its Ultra tation period of 18.5 days for this star, which makes it a rela- Violet spectrum. For instance Butler et al. (1981), Linsky tively slow rotator. Gl 664 was also detected in X-rays, which et al. (1982), Ayres et al. (1983), Oranje (1986), Butler et al. (1987), Doyle (1987), Byrne & Doyle (1989, 1990), Doyle is evidence of a hot corona, with an X-ray luminosity Lx = 5.6 × 1027 erg s−1 (Hunsch et al. 1999). et al. (1990), Landsman & Simon (1993) and Quin et al. (1993) have investigated IUE observations. More recently, Maran et al. Hα is the strongest for gl 664 in our sample stars with an EW (1994), Linsky & Wood (1994), Wood et al. (1997), Redfield & of 640 ± 50 mÅ. Its CaII H & K fluxes are relatively strong, in- Linsky (2002) studied HST/GHRS spectra of AU Mic. Pagano dicating that it is an intermediate activity star. Its strong Hα and et al. (2000) have investigated HST/STIS spectra. Del Zanna Hβ also reflects the fact that Balmer lines increase in absorption et al. (2002) analyzed FUSE observations while Schrijver et al. with decreasing spectral type. Pa was detected but is relatively (1995), Monsignori Fossi & Landini (1994), and Monsignori weak with an EW of 84±10 mÅ. Unfortunately there are no grids Fossi et al. (1996), report on observations at shorter wavelengths of model chromospheres available for dK5 stars. It is therefore with EUVE. Most of these studies report spectral line fluxes difficult to define its activity level. The comparison of these ob- of UV lines and related empirical correlations. Linsky & Wood servations with the models of dM1 stars of Paper VI indicate that (1994) found that C IV and Si IV line profiles show two compo- −1 −1 this is an intermediate activity star, with an Hα absorption close nents with respective widths of about 29 km s and 170 km s . to the maximum value for dK5 stars. Wood et al. (1997) interpret this line broadening in terms of E. R. Houdebine et al.: Observations of M and K dwarf chromospheric lines 939

Fig. 10. We show here the difference spectra AU Mic-Gl 887. Note the wings for the 8542 Å line.

Fig. 11. We show here the difference spectra AU Mic-Gl 908. possible heating mechanisms. Pagano et al. (2000) also report on equivalent widths of –2.01 Å and –2.08 Å. Doyle et al. (1997) −11 −2 −1 two-Gaussian fits to the line profiles. Emission measure distri- inferred a Lyα flux of 1.1 × 10 erg cm s at Earth, which butions were derived notably by Butler et al. (1987), Quin et al. yield a Lyα to Hα ratio of 3.3. Short & Doyle (1998) report on (1993), Maran et al. (1994), Schrijver et al. (1995), Monsignori high resolution observations of the Paβ line and proposed semi- Fossi et al. (1996), Pagano et al. (2000). Del Zanna et al. (2002) empirical model atmospheres to reproduce their observations. derived a rather complete and interesting emission measure dis- This line appears to be in absorption which gives interesting con- tribution from 104 Kto2× 107 K from combined observations straints on model atmospheres, especially about the lower chro- of FUSE, STIS and EUVE spectra. mosphere where the Paschen lines are formed (Short & Doyle 1998, Papers II and III). Worden et al. (1981) also observed the Robinson et al. (1990) measured chromospheric fluxes in Nai doublet for which they detected a core emission. the Caii H and K lines of 4.99 × 105 erg cm−2 s−1 and 5.66 × 5 −2 −1 10 erg cm s respectively. Mathioudakis & Doyle (1989) Pettersen (1989) has observed a number of chromospheric ii = . −2 −1 measured an Mg surface flux of logMgII 6 06 erg cm s . lines of interest for AU Mic. For the Na i D1 and Na i D2 line For Hα, we measured an EW of −2.47 ± 0.1 Å. Robinson et al. core emissions he gives an equivalent width of –0.05 Å and (1990) measured an equivalent width of –2.01 Å. Herbst & –0.10 Å and FWHM of 0.22 Å and 0.40 Å respectively. Here, Layden (1987) derived photometrically an Hα equivalent width we measured line core equivalent widths of −0.063 ± 0.01 Å and of –1.90 Å. Worden et al. (1981) also obtained high resolu- −0.032 ± 0.01 Å for the D1 and D2 lines respectively. Pettersen tion observations of the Balmer lines. They measured equiva- (1989) also obtained unique observations of the Ca ii infrared lent widths of –8.70 Å and –1.56 Å for Hα. The former value, triplet: he found these lines to have an emission core with equiv- which is very high, probably corresponds to a flare. Similarly, alent widths of –0.42 Å, –0.55 Å and –0.37 Å respectively for for Hβ they measured a high value of –3.68 Å, which is much the 8498 Å, 8542 Å and 8662 Å lines. However, some of these larger than our value of –1.9 Å. Pettersen (1989) measured Hα observations have a rather low signal to noise and the spectral 940 E. R. Houdebine et al.: Observations of M and K dwarf chromospheric lines resolution was typically 0.46 Å at Hα, giving a resolving power classify this star as a flare star. For AX Mic, we observed the of only 15 000. Our observations have a higher resolving power H ,H,H,Caii H, Ca i 6572 lines, the Ca ii IR 8498 Å and ii α β  and better signal to noise ratio. For the Ca 8498 Å and 8542 Å 8542 Å triplet lines, the Na i doublet and the Pa lines. − ± line cores, we measured equivalent widths of 0.38 0.05 Å, Our observations of Hα show that this star has a large EW of −0.45 ± 0.05 Å respectively. These values are a bit weaker than 0.592 ± 0.05 Å (Robinson et al. 1990,give0.60Å).ThisEW is those of Pettersen (1989). partly due to the relatively early spectral type of AX Mic. Our Giampapa et al. (1982) derived a 0.54 and 0.42 plage filling observations of the CaII H line shows that this line is weak com- factor based on Caii and Mgii observations respectively, whereas pared to most of the equivalent widths observed for intermedi- in Paper II we derived a 0.30 filling factor based on the modelling ate activity dM1 stars (Paper XI); the EW of this line is only of the Balmer and Lyman lines. More recently, in Paper XII, we –0.599 Å. Similarly, Robinson et al. (1990) give a surface flux 2 derived a plage filling factor of only 9% from a two-component for the CaII H line of only log FH = 4.52 erg/s/cm and Panagi 2 model chromosphere. These values represent a large scatter: un- & Mathioudakis (1993)givelogFH = 5.05 erg/s/cm . Jenkins ffi 2 fortunately, at this stage it is di cult to reach a conclusion. In et al. (2006) found a similar value of log FH = 4.93 erg/s/cm . Paper XII we find that a new grid of model chromospheres with Only the detailed NLTE modelling can explain why we have a lower temperature minima would be preferable. But the two- strong Hα and weak CaII fluxes. But the early comparison with component chromospheric modelling seems the most accurate the models for the near spectral type dM1 stars (Papers XI and approach, particularly since the quiescent region contribution is XII) shows that this is a relatively low activity star, and its large important (Paper XII). Hα equivalent width is at least partly due to its earlier spectral Similarly to Pettersen (1989) we observe a core emission in type. the NaI doublet lines as well as in the CaII IR triplet lines. This The CaII H line has an interesting profile; it shows a central core emission in the CaII IR triplet is pretty strong. The Pa line absorption indicating a relatively high chromospheric opacity. is detected (Fig. 8). It’s profile is interesting: there seems to be This is in agreement with the fact that stars with large radii are wing emission and a core absorption. This profile is broader than slightly metal rich (Houdebine 2008, Paper VII). This can also for other stars, probably because of rotation. This spectral line be due to the fact that in dK stars the chromospheric temperature is of great interest for chromospheric modelling because it is gradient is smaller than in dM stars, simply because the tempera- formed at a lower temperature than the Balmer lines (Short & ture minimum is higher. This smaller temperature gradient yields Doyle 1998a; Papers II and III). a higher mass loading in the region of formation of the CaII lines, We also observed the Pa8 line for AU Mic. Unfortunately we and therefore a higher optical depth for a given column mass of did not observe a reference star for this line, but we observed it the transition region. Calculations by Houdebine (1994) clearly also for AX Mic. If we compute the difference of these two spec- show that the relative intensity of the central absorption in the ff tra AU Mic-AX Mic, the Pa8 line should appear as a P-Cygni CaII lines increases with increasing e ective temperature for a profile: an emission component from AX Mic and an absorp- given transition region column mass. tion component for AU Mic, because they do not have the same The log FMgII of 5.18 given by Mathioudakis & Doyle (1989) . The reason for calculating such difference spec- (Panagi & Mathioudakis 1993, give 5.29) is typically that of a tra is to remove the telluric lines which are strong at this wave- dK7 star. Byrne & Doyle (1989) give the surface fluxes in some length (Fig. 9 for Pa8 and Fig. 8 for Pa). In the case of AU UV lines for AX Mic, but these values are somewhat erroneous Mic-AX Mic, all photospheric line profiles appear as P-Cygni because they use a stellar radius of only 0.58 R. Hunsch et al. profiles because these two stars have close spectral types. (1999) found a very small X-ray luminosity of 1.1 × 1027 erg/s. We computed the difference spectrum; AU Mic-AX Mic. We This is in agreement with the small CaII H flux. show the difference spectrum in Fig. 9. We can note a number of We detected Hβ and Pa in absorption in AX Mic. On the spectral features, but none at the expected wavelength of the Pa8 other hand, Pa8 was not detected (Fig. 9). line (9015 Å). Therefore this line is not detected either for AU Mic nor AX Mic. 3.8. Gl 884 Is Gl 729 more active than AU Mic? Our observations show that there are interesting and important differences between the Gl 884 was often classified as a dM1 star (e.g. Robinson et al. − = two stars. A weak core emission in the NaI doublet lines is de- 1990), however, its infrared color (R I)c 0.770 (Robinson tected for both stars (Fig. 3). But, contrary to AU Mic, there is et al. 1990) is not compatible with this spectral type (Paper VI). no core emission in the CaII IR triplet lines for Gl 729 (Figs. 5 In The CDS database we find a spectral type of dK5 which is in to 7). In addition Hα is much stronger in AU Mic than in Gl 729, better agreement. More recently, Gray et al. (2006) studied this but the surface fluxes in the Ca ii resonance lines are larger for star and found a spectral type of dK7. For Gl 884, we observed Gl 729 than for AU Mic. At this stage we cannot explain the dif- the Hα,H,Caii H, Ca i 6572 lines, the Ca ii IR 8498 Å triplet ferent properties of the chromosphere for these two stars, all we line and the Pa line. can say is that there are important differences. The detailed mod- Robinson et al. (1990) found a surface flux in the CaII H & 2 elling of these lines will have to explain what these differences K lines of respectively log FH = 4.88 erg/s/cm and log FK = are due to. 4.93 erg/s/cm2,andanEW of the CaII K line of –1.32 Å. Here, we find an EW of the CaII H line of −0.97 ± 0.05 Å. Rauscher &Marcy(2006) found mean EW of –0.79 Å and –1.14 Å re- 3.7. AX Mic (Gl 825) spectively for the H and K lines. Gray et al. (2006) list a CaII = AX Mic is a bright (v = 6.64) dK7 star according to its molecu- H & K flux of log FCaII 4.643. Wright et al. (2004) observed lar bandstrengths (Hawley et al. 1996). This star has often been a CaII S index of 1.569. Hα is relatively strong with an EW of classified as a dM1 star (e.g. Byrne & Doyle 1990; Jenkins et al. 0.605 ± 0.05 Å. Herbst & Layden (1987) found an EW of 0.53 Å 2006), but its infrared color (R − I)c = 0.865 is too small for and Panagi & Mathioudakis (1993) quote an EW of 0.60 Å in a dM1 star (typically 1.2, Paper VI). Gershberg et al. (1999) better agreement with our measurement. E. R. Houdebine et al.: Observations of M and K dwarf chromospheric lines 941

26 Although the CaII line surface flux is relatively low, this is luminosity; Lx = 7 × 10 erg/s, and with the fact that Gl 887 one of the most active stars in our sample, according to its CaII is slightly metal poor dM1 star (Paper VII). H line EW (Fig. 1). Similarly to AX Mic, the CaII H line shows a We subtracted the spectra of Gl 887 from those of AU Mic. weak central absorption. H is just about detected in weak emis- We show the difference spectra in Fig. 10. The CaII H line sion although Hα is in absorption. The CaII IR 8498 Å line is does not show any central absorption, in spite of the fact, that rather broad and it seems there is a “filling-in” of the line core for dM1e stars the chromosphere is optically thick in this line when compared to the spectrum of AX Mic for instance (which (Paper VI). There are interesting differences between the CaII is deeper). This would be consistent with the much larger CaII H IR triplet lines at 8498 Å and 8542 Å. The latter line is a bit EW for Gl 884. This could be an interesting constraint for mod- broader and shows some weak wings while the 8498 Å line does elling. It is interesting to note that there is a possible weak filling not. These are interesting constraints for modelling. in for this line for the two dK stars, Gl 664 and Gl 884, whereas we do not observe such filling in for the supposedly more active dM4.5e star Gl 729. 3.10. Gl 908 The Pa line was detected in absorption and is the weakest  − = in our sample, with only 46 ± 10 mÅ, that is almost four times Herbst & Miller (1989) gave an infrared color (R I)K 0.87 − = weaker than in AU Mic. (i.e. (R I)C 1.124). Therefore, this is a dM1 star that suits the selection criterion of Paper VI. It has the same spectral type as AU Mic and Gl 887. In Paper VI, we found that for Gl908: 3.9. Gl 887 R = 0.498 R and [M/H] = −0.271. Bonfils et al. (2005)give an even lower metallicity of [M/H] = −0.52. Therefore, Gl Robinson et al. (1990) gave an infrared color of (R − I) = 1.04 c 908 is a small radius (compared to solar metallicity M1 dwarfs, for Gl 887. This is very close to the infrared color of AU Mic Paper VII) and metal poor M1 dwarf, that probably has a very (R − I) = 1.10 (Paper VII). It is in fact preferable to compare c low activity level (Paper VI). For Gl 908, we observed the H , the R − I infrared colors which are a good indicator of the ef- α H ,Caii H, Ca i 6572 lines, the Ca ii IR 8498 Å and 8542 Å fective temperature rather than the spectral type which varies for  i different classifications for M dwarfs (Paper VI). Gl 887 was triplet lines and the Na doublet. classifiedasanM0.5star(SIMBAD-CDS). It has therefore also Stauffer & Hartmann (1986) found an Hα EW of 0.40 Å. a spectral type close to that of AU Mic. Hence, we can compare Herbst & Miller (1989) found 0.44 Å, Herbst & Layden (1987) their spectra and look for the real contributions of plages in AU found 0.36 Å, and here we also find 0.36±0.05 Å. This is a weak Mic. For Gl 887, we observed the H,Caii H lines, the Ca ii IR Hα EW according to Paper VI. This agrees with the small CaII 8498 Å and 8542 Å triplet lines and the Pa line. H EW of only −0.188 ± 0.03 Å that we observe. The average Lopez-Morales (2007)gaveR = 0.491 ± 0.014 R (see also CaII H & K equivalent width from various authors is –0.245 Å Segransan et al. 2003)and[M/H] = −0.22 ± 0.09 dex (see also (Paper VI). The average of the two lines is a better diagnostic Bonfils et al. 2005). This can be compared to the values found in of activity because the line ratios vary significantly with activ- Paper VII from the radius/metallicity correlation; R = 0.579 R ity level (Paper VI). This is a low metallicity, very low activ- and [M/H] = −0.092 dex respectively. The values of radii in- ity dM1 star (Papers VI and VII). The surface flux in the CaII ferred in Paper VII agree well in general with interferometric H & K lines is log F = 4.88 erg/s/cm2. Byrne (1993)gavean ff 2 measurements. The di erence in radii between the measure- even lower value of log F = 4.56 erg/s/cm . LX given by Hunsch ments of Lopez-Morales (2007) and ours is one of the largest et al. (1999)is1.7 × 1027 erg/s. This is also in agreement with in our sample in Paper VII. At this stage we have no explanation its low vsini (0.9 km s−1, Paper VII) and long rotation period for this relatively important difference (error in , in R − I P/ sin i = 28 d. or in interferometry?). Additional measures will be necessary to We can therefore compare the Hα and NaI D1/D2 spectra reach a conclusion. The metallicity we derived in Paper VII for of AU Mic and Gl 908. We cross-correlated the spectra of AU Gl 887 directly depends on its radius. This is the reason why ff Mic and Gl 908 for a region where narrow photospheric features there is also some di erence between the two metallicity mea- are present (not much affected by metallicity) in order to deter- surements. Therefore, Gl 887 has a small radius (compared to mine the wavelength shift. We shifted the spectrum of AU Mic solar metallicity M1 dwarfs, Paper VII) and is a metal-poor dM1 to the reference frame of Gl 908. We show the difference spec- star. tra in Fig. 11.TheH profile no longer shows a self-reversal − ± α Gl 887 has a rather weak CaII H line EW ( 0.796 0.05 Å) and is clearly asymmetric in the line centre. We can also note − ± compared to AU Mic ( 4.55 0.2Å).H also seem to have some wing emission. In Paper II, we showed that the central ab- − ± a weak emission in Gl 887 ( 0.11 0.03 Å), whereas Hα is sorption of the Hα line is the main constraint on the models. in absorption (EW = 0.45 Å, Panagi & Mathioudakis 1993). A weaker absorption implies a higher transition region pressure Rauscher & Marcy (2006) give equivalent widths of –0.77 Å and and a higher chromospheric gradient. Here we did not take into –0.54 Å respectively for the H and K lines, from several mea- account a filling factor for the quiescent regions, which means surements. This is close to the value we observe. Robinson et al. that in reality the central absorption is a bit larger than in Fig. 11. 2 (1990)giveCaIIH&KfluxesoflogFH = 4.68 erg/s/cm and However, we conclude that our model given in Paper II is still a 2 log FK = 4.62 erg/s/cm respectively. Tinney et al. (2002)gave lower limit of the true plage model of AU Mic.  = / / 2 log RHK 5.01 erg s cm (see also Jenkins et al. 2006). Doyle Figure 11 also shows that the NaI D1/D2 lines have inner 2 et al. (1990) found Lyα and MgII surface fluxes of 5.32 erg/s/cm wing emission as well as the core emission. In fact, most of and5.23erg/s/cm2 respectively. The MgII flux is substantially the energy is contained in these wings. Again, these are inter- larger than the CaII flux. esting constraints for the modelling of dMe chromospheres. The According to the CaII and Hα equivalent widths, and to the narrow absorption features visible on the red wings of the NaI models of Houdebine & Stempels (1987, Paper VI), Gl 887 D1/D2 lines come from the low pressure sodium lamps of the is a low activity dM1 star. This agrees with its low coronal nearby city of La Serena visible in the Gl 908 spectrum (Fig. 3). 942 E. R. Houdebine et al.: Observations of M and K dwarf chromospheric lines

3.11. McC 332 (HIP 104383)

McC 332 is one of the most interesting stars that we observed. It was reported as a binary by Poveda et al. (1994). Later, speckle interferometry was reported by Balega et al. (2004). The pri- mary has been classified as a dM2 star with v = 9.44, while the secondary is a faint M dwarf with v = 13.4. The primary has (R − I)c = 0.68 and the secondary has (R − I)c = 1.42 (Weis 1993). The primary’s color is too blue for an M2 star, it is more probably a dK7 star (Bessel 1990). The color of the secondary indicates that it is a dM3.5e star (Bessel 1990). Gershberg et al. (1999) reported McC 332 as a flare star. It is most probably the fainter active companion that exhibits flares, because as shown from the spectra below it is much more active than the primary (Hα emission, stronger Ca ii H emission). For McC 332 we ob- served the Hα,Hβ,H,Caii H, Ca i 6572 lines, the Ca ii IR 8498 Å and 8542 Å triplet lines and the Na i doublet. McC 332 has been rarely observed. Wilson and Wooley (1970) reported a CaII H & K index of +3, which means a rela- tively strong emission. Herbst & Layden (1987) observed an Hα EW of 0.63 Å photometrically. Here we report a much weaker Hα EW, of only 0.29 Å. Our Hα line profile is of P-Cygni type (Fig. 4) and our CaII H profile clearly shows a blue asymmetry Fig. 12. We show here the cross-correlation profiles of Gl 488 with it- (Fig. 14)! This is the first time that this star has been observed self (dotted line) and of McC 332 with Gl 488 (dashed line). The two spectroscopically. Herbst & Layden (1987) probably only ob- Gaussian fit of the latter profile is also shown. served the primary, whereas here we have spectra of both the primary and the secondary. Table 2. Summary of the two Gaussian fits to the correlation profile and the CaII H line. We extracted the 6565–6590 Å region of our Hα spectra of McC 332 and Gl 488. This region is interesting for cross- correlations because it includes strong and unblended absorption Profile λ0 FWHM Max λ0 FWHM Max features. We cross-correlated Gl 488 with itself and McC 332 Type (Å) (Å) flux (Å) (Å) flux with Gl 488. The resulting normalized cross-correlation profiles Correlation –1.27 0.395 0.880 –0.888 0.314 0.17 are shown in Fig. 12. We chose Gl 488 instead of Gl 908 which CaII H 3967.04 0.24 0.879 3967.26 0.23 2.96 has a low vsini (Houdebine 2008) because the spectrum has a higher S/N ratio. The cross-correlation profile of Gl 488 has a FWHM of the finding that McC 332 is a binary. The wavelength shift yields a velocity of 9 km s−1. 0.347 Å. This is 0.19 Å larger than the expected instrumental The CaII H spectrum of McC 332 is shown in Fig. 14.We FWHM for a resolution of 40 000. This yields a possible vsini of − can see a clear asymmetry in this profile. The profile fitted by 8.2 km s 1 for Gl 488 which seems too high for an inactive dK7 −1 two Gaussians yields a wavelength difference of 0.22 Å, i.e., star (typically less than 3 km s see Paper VII). Therefore, either −1 the resolution was lower, or the width is due to the intrinsic width 16.6 km s . Note that the Hα and CaII H lines were not ob- of the line profiles. All we can say is that the McC 332/Gl 488 served at the same time. We obtain a mean redshift for the sec- ondary component of about 14 km s−1. Note that the CaII H line cross-correlation profile is broader with a FWHM of 0.394 Å. − is stronger for the secondary than for the primary. The parame- This means that vsini for McC 332 is about 8.5 km s 1 larger ters of our two-Gaussian fit are summarized in Table 2. than vsini for Gl 488. This is quite large for an Hα absorption line star (Paper VII). This excess broadening is probably partly due to the fainter secondary Gaussian component discussed be- 4. Conclusion low. We have observed a number of spectral lines for a few dK, We can note in Fig. 12 that the McC 332/Gl 488 cross- dM and dMe stars that provide new constraints for the purpose correlation profile is asymmetric and shows a bump in its red of NLTE-radiation transfer modelling of their chromospheres. wing. We fitted two Gaussians and found that the second fainter Notably, Pa was observed for the first time. We plan to use these Gaussian is red shifted by 0.36 Å with respect to the main data to build new model chromospheres. The Paschen lines are Gaussian. We believe this is evidence that McC 332 is a spectro- complementary to the Balmer lines as they are formed lower in scopic binary with a dK7 principal component and a secondary the chromosphere (Short & Doyle 1998a). It would be interest- component of later spectral type. The wavelength shift yields a −1 ing to reconcile the modelling of these lines, especially since velocity of 16 km s only. The parameters of our two-Gaussian Short & Doyle (1998a) could not model both the Paschen and fit are summarized in Table 2. Balmer lines at the same time. With our detection of Pa we also We subtracted the Gl 488 Hα spectrum from the MCC 332 have better constraints than those of Short & Doyle (1998a)who ff Hα spectrum. The di erence spectrum is shown in Fig. 13.A barely detected the Paβ line. nice emission Hα profile appears with a possible self-reversal. We have gathered complete sets of observations of This emission profile is red-shifted by about 0.20 Å with respect chromospheric lines from the lower chromosphere to the tran- to the Hα absorption profile of Gl 488. This further strengthens sition region. These sets are unique and are precious to develop E. R. Houdebine et al.: Observations of M and K dwarf chromospheric lines 943

We have observed a bright spectroscopic binary, composed of a dK7 primary and a dM3.5e secondary. The secondary is quite active with strong Hα and CaII H emissions. We plan to gather more such high resolution observations of K and M dwarfs in order to build more accurate grids of model chromo- spheres and for different spectral types from dK5 to dM8.

Acknowledgements. The authors would like to thank Dr. C. J. Butler from Armagh Observatory for improving the English of this manuscript.

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